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Lives of the Stars Lecture 6:
Stellar evolution
Presented by
Dr Helen Johnston
School of Physics
Spring 2016
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In tonight’s lecture
• Changes on the main sequence
– gradual changes during the star’s life
• The beginning of the end
– what happens when the star runs out of hydrogen
• Different fates
– how low-mass and high-mass stars end differently
• The evidence for stellar evolution
– star clusters
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A lingering death
You recall from our discussion of stellar structure that main sequence
stars are in a dynamic equilibrium: they are winning the battle against
gravity, but only at the cost of burning fuel: hydrogen.
What happens when the star runs out of fuel? That leads to massive
changes in the star’s structure and appearance: stellar evolution.
Tonight we will look at the profound changes a star goes through
during the last stages of its life.
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The fact that gravity determines the structure of the star explains why it
is the star’s mass which is by far the most important variable in what it
looks like.
104
L (Lsun)
This is exemplified in the
Herztsprung-Russell diagram, the
most obvious feature of which – the
main sequence – is a mass sequence,
not a time sequence. An F-type star
formed as an F-type star, and will
remain an F-type star during the
entire main sequence stage.
106
102
high mass
stars
1
10–2
low mass
stars
10–4
40000
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main sequence
20000
10000
T(K)
5000
2500
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Changes on the main sequence
Stars do change somewhat while they are on the main sequence.
While a star is on the main sequence, it is burning hydrogen in its core.
Helium is formed, and gradually builds up as a sort of ash: helium
poisoning.
Helium can’t burn, because it needs a temperature of about 100
million degrees before it can ignite. So it just builds up in the core.
helium builds up
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As each fusion converts four H atoms to one He atom, the core of the star
has fewer atoms in it, so the pressure goes down. Gravity squeezes the
core more tightly, which increase the temperature, which increase the
rate of fusion. This produces more energy, which makes the outside
layers of the star expand a bit, which makes the star a bit brighter and
a bit cooler.
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These changes in the Sun will have a profound effect on the Earth. When
the Sun began its main sequence life about 5 billion years ago, it was
only 70% as bright as it is now. In another 5 billion years, it will be
roughly twice as bright, which will raise the average temperature of the
Earth at least 19° C.
In fact, we have a hard time
reconciling these predicted
temperatures with the geological
record (the “faint young sun
paradox”).
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How long a star can burn hydrogen on the main sequence depends on
two things:
• how much hydrogen it has; and
• how fast it burns it
The first is just the star’s mass, and the second the star’s luminosity.
But we know how those two are related: we derived the mass-luminosity
relation when we were looking at stellar models in Lecture 3.
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how fast you
burn it
how much fuel
Here it is again: as we increase the mass of the star, the luminosity increases
enormously: a factor of 10 increase in mass corresponds to a factor of 3000
increase in luminosity.
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We can write this dependence of the star’s luminosity on its mass as
L ∝ M3.5
so the lifetime of a star goes like
t ∝ M/L ∝ M/M3.5 = M–2.5
In other words, a factor of 10 increase in mass corresponds to a
decrease in the lifetime of the star by a factor of 300.
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So we can make a table showing the lifetime for stars of different
masses:
Spectral
type
O5
B0
A0
F0
G0
K0
M0
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Mass
(Sun=1)
Luminosity
Years on main
sequence
40
15
3.5
1.7
1.1
0.8
0.5
405,000
13,000
80
5.4
1.4
0.46
0.08
1x106
11x106
440x106
3x109
8x109
17x109
56x109
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Interestingly, the amount of energy released per kilogram is almost
identical for all types of stars.
energy per kilogram = L t / M
∝ M3.5 x M–2.5/M
= constant
Massive stars are much more luminous, but live for much less time, so the
amount of energy extracted per kilogram is identical to the faintest M
stars.
This is, of course, because all stars are getting their energy from the
same source: nuclear reactions.
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The beginning of the end
When the core runs out of hydrogen, fusion stops so the core contracts
and heats up. This heats the unprocessed hydrogen just outside the
core, which then ignites in a spherical shell, burning outwards. This
sudden increase of energy forces the outer layers of the star to expand dramatically, so the star becomes enormously large and red: a red giant. Stars like the Sun can increase in radius by a factor of 10–100, while more massive stars can increase by a factor of 1000.
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On the HR diagram, the star moves dramatically upwards (getting much
brighter) and somewhat to the right (getting cooler). This is called the
giant branch.
For the first time, the star changes its spectral type. A G-type star
may end up as a high-K or low-M giant.
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Although the H-fusion shell can force the envelope to expand, it cannot
stop the collapse of the core. This continues to collapse, until the inner
12% of the mass is compressed to very high density. At such extreme
densities, the gas stops obeying the ideal gas laws.
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The collapse of the core is stopped by electron degeneracy, which might
better be called quantum pressure. Electrons in the core can only have
certain energies, and quantum mechanics states that two electrons
cannot occupy the same energy level. As the gas becomes very dense,
all the available energy levels are occupied. Since the electrons cannot
change their energy, the gas cannot be compressed, so it begins to
behave like hardened steel.
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Low-density gas
(non-degenerate)
High-density gas
(degenerate)
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This enables the core to resist the inward pressure of the outer layers of
the star.
However, now that the gas can no longer be squeezed, the star has lost
its thermostat. Recall that a star on the main sequence regulates its
energy production by way of the ideal gas laws:
core loses energy ➔ pressure drops
➔ star contracts
➔ temperature increases
➔ energy production increases
➔ equilibrium is restored
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In a degenerate gas, however, pressure does not depend on
temperature. If a star supported by degeneracy pressure loses energy,
the pressure remains unchanged. The star does not contract, and there is
no increase in temperature to counteract the loss of energy.
Similarly, if the star heats up, the pressure does not increase so as to
expand the core and lower the temperature, so the temperature can
keep growing.
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The helium flash
The core continues to contract and the temperature keeps rising. Once
it reaches 100 million K, helium can ignite to form carbon. Two nuclei
temporarily form an unstable atom of beryllium; if a third helium
nucleus reacts with beryllium, it forms a carbon atom.
In low mass stars (up to 3 solar masses), the turn-on of helium fusion is
very sudden: the helium flash.
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Because the core is degenerate, when helium ignites and produces more
energy, the temperature of the core rises but the core does not expand.
This means that the rate of energy production keeps going up, which
keeps increasing the temperature, in a runaway process. In a few hours
the temperature jumps to a billion degrees, producing more energy
than the entire galaxy.
But this energy doesn’t escape the star: it all goes into removing the
electron degeneracy. Now the core can behave like a perfect gas
again, so the star’s thermostat has been restored: the core can expand
and cool.
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Quickly the P-T thermostat brings the reaction under control, and the
star begins fusing helium steadily in the core.
More massive stars don’t have a helium flash, because their cores are not degenerate when helium fusion begins, so the runaway process never takes place: helium fusion begins gradually.
Core close-up
Actual size
1/1000 of star
Envelope
Helium core
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Burning shell
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Aside: Why does He fusion need high temperatures?
Helium nuclei contain four positive charges, so in order to fuse
together they have to overcome a much higher electromagnetic
(Coulomb) barrier than two hydrogen nuclei (one positive charge
each). Fusing heavier elements takes even higher temperatures.
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4H → He
5 million K
3He → C
100 million K
C + He → O
600 million K
2C → Mg
1,000 million K
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The star is now making energy in two places at the same time. Only in
the core is it hot enough to burn helium, but just outside the core it’s now
hot enough to burn hydrogen for the first time. The contraction of the
core stops, and the envelope of the star contracts and grows hotter.
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Helium-burning
core
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This makes the star move down and to the left in the HR diagram: the
horizontal branch.
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The cycle then repeats itself. Helium is fused to carbon and oxygen in
the star’s core, and carbon and oxygen ash build up. Eventually the core
runs out of helium, so fusion stops, and the core resumes its collapse. This heats up the layer of helium just outside the core, which begins fusing in a shell of its own, inside the hydrogen burning shell.
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This extra energy forces the envelope to expand again, and the star
begins a second ascent of the giant branch: the asymptotic giant branch.
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We can predict what happens next: the cycle repeats again and again,
each time starting to fuse heavier and heavier elements to stave off the
next collapse. The core reaches high temperatures by converting
gravitational energy into thermal energy.
However, if the star is not massive enough, it will never reach high
enough temperatures to fuse carbon: this is true for the Sun.
So for low mass stars, this is the end of the line. We will now look at the
fate of different mass stars: very low mass (< 0.4 solar masses),
medium mass (0.4 – about 6 solar masses) and high mass (> 6 solar
masses).
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Very low mass stars
Very low mass stars (mass less than about 0.4 solar masses) are
different in one important respect from heavier stars: their interiors
are fully convective. The fused helium is stirred through the whole star,
so it has the whole of its hydrogen mass to prolong its stay on the main
sequence.
Because they don’t have a core, very low mass stars can never
develop an inert helium core surrounded by unburnt hydrogen, so can
never become a red giant. Instead, when the star has used up
essentially all their hydrogen, fusion will stop, the star will contract and
heat up, then gradually cool as a helium white dwarf.
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Medium mass stars
Intermediate mass stars, with masses between 0.4 and about 6 solar
masses, have another evolutionary path. These stars get as far as
burning helium in their cores, but cannot get hot enough to ignite
carbon.
Thus when the helium core is exhausted, they cannot resist the
gravitational force, and the core finally collapses to form a white
dwarf.
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The evolution of a solar-like
star is summarised in this
picture.
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However, in the final supergiant stage, the surface gravity is so low that,
during the collapse, the star loses its outer layers in spectacular fashion.
The entire outer envelope is blown off
in a massive stellar wind, in a series of
thermal pulses. As the envelope of the
star falls towards the contracting core,
some of the material ignites, causing
the outer surfaces to bounce and
vibrate. Eventually the outer 40% of
the star’s mass will be "coughed" into
space, floating outward in a concentric
set of spherical bubbles.
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The core, now exposed, continues to contract and heat: in only a few
thousand years it has reached a temperature of 30,000 K, so it
becomes very blue.
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Naively, we might expect
the gas to be ejected
spherically; the nebula,
however, would appear as a
ring because we see more
gas along the sides of the
bubble. This is the Helix
nebula (NGC 7293), the
nearest planetary nebula.
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This animation shows the birth of the Helix Nebula. The star ejects its
outer envelope at low velocity, exposing the hot core of the star. This
star has a fast low-density wind that blows a big cavity in the dispersed
envelope. UV radiation from the star heats the gas, causing it to glow.
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However, many planetary nebulae
have strange and wonderful
shapes, indicating that the mass
ejection process is very
complicated. In binary systems, for
instance, the gas outflow is
influenced by the interactions
between the stars. This is an HST/
X-ray picture of the Cat’s Eye
nebula (NGC 6543).
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The Butterfly nebula (M2-9) is very far from circular. The central object
is a binary orbiting within a gaseous disk; the expelled envelope
interacted with this disk to form the spectacular shape we see today. In
fact, most planetary nebulae are bipolar.
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Simulations suggest that the strange shapes of planetary might be due
to asymmetric winds from the progenitor confining the gas ejected
during the planetary nebula phase. A disk-shaped cloud of gas around
the star’s equator forces the gas to escape towards the poles, forming a
bipolar nebula.
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Simulation of gas escaping from a star, while being constrained by a
torus of gas around the star.
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These simulations reproduce
some of the known planetary
nebulae with amazing accuracy.
Hubble image of the “Hourglass nebula”,
MyCn 18, compared to the interacting
wind model viewed from an angle of 40o.
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The “Ant Nebula”, Menzel 3.
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M27: The Dumbbell Nebula
The remaining core of the star, consisting mostly of carbon and oxygen
(ash from the helium burning), starts off very hot, but has no fuel source,
i.e. no nuclear reactions are taking place to replace the energy lost by
radiation. It will only cool and shrink: it has become a white dwarf.
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The planetary nebula NGC 2440,
with its central white dwarf, one of
the hottest white dwarfs known.
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The planetary nebula known as the “Skull
Nebula”, NGC 246, with its white dwarf,
the fainter member of the binary star
The
system
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Sydney at the nebula's centre.
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Massive stars
High mass stars (greater than about 6 solar masses) have much more
violent endings.
Massive stars can reach
temperatures in their core high
enough to begin fusing carbon,
producing neon and oxygen. When
the carbon is exhausted in the core,
it contracts and carbon ignites in a
shell. This pattern of core ignition
and shell ignition continues with fuel
after fuel, until the star develops a
layered structure.
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After carbon fuses, oxygen, neon and
magnesium fuse to make silicon and sulphur,
and then the silicon fuses to make iron.
The fusion of these nuclear fuels goes faster
and faster as the atomic number increases,
both because there are fewer atoms to fuse,
and less energy is released each time.
hydrogen
helium
carbon
oxygen
neon
silicon
iron
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These elements collect and burn in
concentric shells, like an onion.
Once the silicon in the core has fused
to make iron, the star can no longer
support itself against collapse.
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The table shows the time taken for each fusion process for a 25 solar
mass star.
Fuel
Temperature (K)
Time (y)
Sample reaction
Hydrogen
40 million
7,000,000
41H →4He
Helium
200 million
500,000
44He →12C
Carbon
600 million
600
212C →20Ne + 4He
Neon
1.2 billion
1
20Ne
Oxygen
1.5 billion
0.5
216O →28Si + 4He
Silicon
2.7 billion
1 day
228Si →56Fe
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+ 4He →16O + 24He
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Why iron?
Why does the sequence of nuclear burning end with iron? Iron is the
most "tightly bound" of the elements. If an element is to burn by fusion
then energy must be given off when that element is combined with
another. The new nucleus is more tightly bound than the two separate bits; it would take energy to break the new nucleus apart to the original parts.
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But with iron we have the situation where to add any other element costs
energy, that is the new nucleus is less tightly bound. Indeed, up at the
high end of the periodic table we have elements like uranium and
plutonium that give up considerable energy when split apart in the
process of nuclear fission. Iron then represents the energy minimum and
once a star has converted its matter into iron then it has no further
recourse. The star has reached a dead end.
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Abundances of the elements
This sequence explains the first half of
the cosmic abundance diagram: elements
up to iron are formed in stars, and the
relative abundances are well explained
by what we know of fusion processes.
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Nuclear chemist William Harkins (c. 1931) noted that
"elements of low atomic weight are more abundant than those of high
atomic weight and that, on the average, the elements with even atomic
numbers are about ten times more abundant than those with odd
atomic numbers of similar value."
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Hydrogen burning
Helium burning
Carbon burning
Oxygen burning
Silicon burning
But what of elements heavier than iron? How are they formed?
(We’ll find out next week...)
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During the later stages of burning, more and more energy is being lost
to neutrinos: most of the star’s energy, by the time carbon burning starts.
Fuel
Temperature (K)
Time
(y)
Hydrogen
40 million
7,000,000
2.7x1031
–
Helium
200 million
500,000
5.3x1031
<1.0x1029
Carbon
600 million
600
4.3x1031
7.4x1032
Neon
1.2 billion
1
4.4x1031
1.2x1036
Oxygen
1.5 billion
0.5
4.4x1031
7.4x1036
Silicon
2.7 billion
1 day
4.4x1031
3.1x1038
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photon luminosity neutrino luminosity
(J/s)
(J/s)
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While all this fusion is taking place, the star is zig-zagging across the
HR diagram.
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By this stage, the iron core exceeds 1.4 solar masses. When silicon
fusion finishes, the core begins to collapse for the final time. And this
time, there is nothing to stop it. The end result is the destruction of the
star in a supernova explosion: we’ll look at this in detail next week.
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Mass loss
One major complication in understanding the evolution of massive
stars is that they lose large amounts of mass at many stages of their
life.
Hot stars emit a continuous outflow of matter from their surfaces as a
stellar wind. Our own Sun has a solar wind which reaches speeds of
400–700 km/s with a mass loss rate of about 10–14 solar masses per
year. Over a ten billion year lifespan, at this rate the Sun will lose
about 0.01% of its mass to the solar wind.
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By contrast, the winds from hot stars can be a billion times stronger,
losing up to 10–5 solar masses per year at speeds of up to 3000 km/s.
This means that even during the much shorter lives of the stars (a few million years), they can lose on the order of half or more of their mass.
The “Pistol Nebula” and its central star, which may
have weighed up to 200 times the mass of the Sun
The
before
University ofshedding
Sydney
much of its mass in violent eruptions.
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The Bubble Nebula, NGC 7635,
is being pushed out by the
stellar wind of massive central
star BD+602522, which has a
mass about 40 times the mass
ofThethe
Sun. The bubble is about
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3 parsecs (10 ly) across.
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This has substantial implications for the evolution of the star. Reducing
the mass of the star reduces the pressure and temperature in the
interior, which can reduce the mass of the core.
And it is the mass of the core when fusion stops which governs whether
the star explodes as a supernova, and what kind of remnant it leaves
behind.
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Very massive stars, with initial masses greater than 40 solar masses,
become luminous blue variables after leaving the main sequence. During
this phase, which lasts for perhaps 40,000 years, the stars are highly
variable and losing mass through strong winds. At minimum brightness
they appear as blue B-type supergiants; during outburst they are much
redder. Every few centuries they have sudden giant eruptions, ejecting
large amounts of mass.
Radio images of two luminous blue variables, AG Carinae and Henize 3–519,
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showing rings of emission from mass lost during major eruptions.
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Eta Carinae is the most
famous luminous blue
variable. During the 1840s it
brightened by 4 magnitudes,
becoming one of the brightest
stars in the sky. Hubble
images show two huge
bubbles of gas, remnant of
the expulsion of about a solar
mass of material.
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Wolf-Rayet stars represent the
most extreme stage of mass
loss in the life of a massive star:
the star has lost so much mass
that they are actually exposing
the underlying layers which
have already undergone
nuclear fusion.
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HST image of the Wolf-Rayet star WR124,
showing the star surrounded by hot clumps
of gas being ejected at high speed. Page
The amount of mass lost in these stages of a star’s life essentially
determines what its ultimate fate is. We’ll discuss the endpoints of stellar
evolution in a few lectures, but we note here that most of the uncertainty
in what mass of star leads to what kind remnant is because of the
uncertainty in how much mass is lost by the time the star ends its life.
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The final fate of the Earth
Mass loss is also vital for determining the final fate of the Earth. As
the Sun ages, it will expand and become a red giant.
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The final radius the Sun reaches is important. However, as the Sun loses
mass, its gravitational pull on the planets weakens, and their orbits
expand.
Will they expand enough to keep ahead of the expanding photosphere, or will they be engulfed?
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Recent calculations* suggest that although the Sun’s outer surface may
not quite reach Earth, tidal interaction between the Earth and the giant
Sun will drag the Earth inwards, to be engulfed by the Sun.
This will take place just before the Sun reaches the tip of the RGB,
around 7.59 ± 0.05 Gyr from now.
In any case, the Earth is likely to become uninhabitable long before that
point is reached.
* “Distant future of the Sun and Earth revisited” by K-P Schroder & R C. Smith (2008)
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The evidence for stellar evolution
We’ve built up a wonderful picture for stellar evolution: but the
timescales are so long that we can never observe a star evolving. How
certain are we about all this?
The best evidence comes from studies of star clusters. Clusters of stars
are ubiquitous in the Galaxy.
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Generally, we divide clusters
into two classes: open clusters,
consisting of 100–1000
young stars, all born from the
same molecular cloud;
The Jewel Box Cluster (NGC 4755, or
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Kappa
Crucis)
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...and globular clusters, which
are extremely old groups of
10,000–1,000,000 stars,
born in the very earliest
stages of the Galaxy’s
evolution.
The dense globular cluster
The
University
of Sydney 6093)
M80
(NGC
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Since all stars in a cluster form at the same time, they are the same age.
Their HR diagram, therefore, becomes a powerful test of stellar
evolution theory. Stars of different masses, evolving at different rates,
will appear at different places in the diagram.
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This is the HR diagram for the globular cluster M13, an old globular
cluster. We can clearly see the main sequence, the red giant branch,
and the horizontal branch. There is one more notable feature: a gap in
the horizontal branch. Stars are rarely found in this region, and those that are are unstable: they pulsate in brightness.
This region is know as the instability strip.
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The most famous of these variable stars are the Cepheid variables,
which are giants of mass greater than 2 solar masses, which pulsate
with periods of several days. The pressure-temperature thermostat is
out of sync, so the stars expand and contract in a regular fashion. Most
importantly, the higher the luminosity of the star, the longer the period,
so Cepheids can be used as standard candles for determining the scale
of the Galaxy and the Universe.
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Aside: In fact, the more we look, the more it seems that globular clusters
aren’t quite that simple...
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We can summarise the final fates of the different types of stars in the
following diagram.
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Next week
we’ll look in detail at the explosions that end the lives of massive stars:
supernova explosions
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Further reading
• “Cosmic Catastrophes: Exploding Stars, Black Holes, and Mapping the Universe” by J. Craig Wheeler (Cambridge UP,
2007) – the first edition from 2007 was subtitled “Supernovae, Gamma-ray Bursts and Adventures in Hyperspace”–
has a good discussion of how stars evolve; most of the book is taken up with how they die, so I’ll be referring to this book
again in the next couple of lectures. Note: This book is not easy; he covers an awful lot of material. His diagrams are
great, however, and it’s worth sticking with.
• There’s a nice biography of Henrietta Leavitt, the discoverer of the Cepheid period-luminosity relation: “Miss Leavitt’s
Stars: The untold story of the woman who discovered how to measured the universe” by George Johnson (Atlas Books,
2005), which mostly goes to show how little we know of her.
• The Australia Telescope has several very nice pages on their Outreach site. “Post-Main Sequence Stars”, http://outreach.atnf.csiro.au/education/senior/astrophysics/stellarevolution_postmain.html is a good summary of the
material we discussed in this lecture. The “HR Diagram Activities” page http://outreach.atnf.csiro.au/education/senior/astrophysics/stellarevolution_hractivity.html has links to several sites which
allow you to explore stellar evolution in the Hertzsprung-Russell diagram.
• Bruce Balick has a lovely page about planetary nebulae and their formation, at http://www.astro.washington.edu/balick/WFPC2/index.html
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Sources for images used:
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Faint young sun paradox: from http://zebu.uoregon.edu/~imamura/122/lecture-1/lecture-1.html
Mass-luminosity relation: from Auckland Astronomical Society http://www.astronomy.org.nz/aas/MonthlyMeetings/MeetingMay2002.asp
Structure of a red giant: from ASP1022: Lecture Materials http://www.maths.monash.edu.au/asp1022/images/Snow15.16.gif
Triple-alpha process: from Australia Telescope Outreach and Education: Post-main sequence stars http://outreach.atnf.csiro.au/education/senior/astrophysics/stellarevolution_postmain.html
Burning in shells: from Astronomy 122: Birth and Death of Stars: Stellar Evolution from Middle Age to Death, by Jim Brau, http://blueox.oregon.edu/~jimbrau/astr122/Notes/Chapter20.html
The life-cycle of a solar-mass star: from Stellar Structure and Evolution: the evolution of low-mass stars by Vik Dhillon, http://www.shef.ac.uk/physics/people/vdhillon/teaching/phy213/phy213_lowmass.html
Size of red giants: from Wikipedia http://en.wikipedia.org/wiki/Image:Redgiants.jpg
Planetary nebula NGC 2440: Hubble image, from Astronomy Picture of the Day 2004 January 11, http://apod.nasa.gov/apod/ap040111.html
Helix nebula: blend of HST and ground-based images, from NOAO Image Gallery, http://www.noao.edu/image_gallery/html/im0840.html
Animation of the formation of the Helix Nebula: from Hubble Site News Center, http://hubblesite.org/newscenter/newsdesk/archive/releases/1996/13/
Cat’s Eye nebula: combined HST/Chandra image, from Astronomy Picture of the Day 2016 July 24, http://apod.nasa.gov/apod/ap160724.html
Butterfly nebula: HST image, from Astronomy Picture of the Day 2004 February 1, http://apod.nasa.gov/apod/ap040201.html
Interacting wind models from planetary nebulae: from Vincent Icke’s Planetary Nebulae page, http://www.strw.leidenuniv.nl/~icke/html/VincentPN.html
Planetary nebula NGC 2440: from APOD 2005 January 23 http://apod.nasa.gov/apod/ap050123.html
Dumbbell nebula: from APOD 2016 November 2 http://apod.nasa.gov/apod/ap161102.html
NGC 246: from APOD 2006 April 18 http://apod.nasa.gov/apod/ap060418.html
Structure of a massive star: from Astronomy 1301-02 by Tony Hall: Deaths of Massive Stars http://lepus.physics.uair.edu/~tahall/EXAM3/Week3.htm
Binding energy curve: from “General Chemistry” by Hill and Petrucci, Fig. 19.6 http://cwx.prenhall.com/bookbind/pubbooks/hillchem3/medialib/media_portfolio/19.html
Cosmic abundance of the elements: from Astro 105: The Milky Way, Lecture XII http://www.pas.rochester.edu/~afrank/A105/LectureXII/LectureXII.html
Pistol nebula: from Hubblesite http://hubblesite.org/newscenter/newsdesk/archive/releases/1997/33/
Bubble nebula: by Russell Croman, from APOD 2005 November 7 http://apod.nasa.gov/apod/ap051107.html
Radio images of LBVs: from http://www.astro.umd.edu/~white/text/lbv_images.html
Radio images of LBVs: from http://www.astro.umd.edu/~white/text/lbv_images.html
Eta Carinae: from Hubblesite http://hubblesite.org/newscenter/archive/releases/1996/23
WR124: from Hubblesite http://hubblesite.org/newscenter/archive/releases/1998/38/
Sun becoming a red giant: from Science@NASA, Sizzling comets circle a dying star http://science.nasa.gov/headlines/y2001/ast11jul_1.htm
Artist’s impression of planet around a red giant: image by Dirk Terrell http://www.boulder.swri.edu/~terrell/dtart_old.htm
Calculations of the fate of the Earth: “Distant future of the Sun and Earth revisited” by K.-P. Schroder & R. Connon Smith, MNRAS 386 155 http://adsabs.harvard.edu/abs/2008MNRAS.386..155S
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Jewel Box cluster: photo by Dieter Willasch, from Astronomy Picture of the Day 2010 August 17, http://apod.nasa.gov/apod/ap100817.html
Globular cluster M80, Hubble image, from Astronomy Picture of the Day 1999 July 7, http://apod.nasa.gov/apod/ap990707.html
Hertzsprung-Russell diagrams of clusters as they age: from “Explorations: An Introduction to Astronomy” by Thomas Arny, Fig. 13.24 http://www.mhhe.com/physsci/astronomy/arny/instructor/graphics/ch13/1324.html
HR diagram for M13: from Astronomy 122: Birth and Death of Stars by Jim Schombert, http://abyss.uoregon.edu/~js/ast122/lectures/lec15.html
Multiple populations in omega Cen: from “The quite complex “Simple Stellar Populations” of globular clusters” by Angela Bragaglia, http://arxiv.org/abs/0912.5280
Supernova image: SN 1994D from Hubble, http://www.spacetelescope.org/images/opo9919i/
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